Astronomy
&
Astrophysics
A&A 557, A23 (2013)
DOI: 10.1051/0004-6361/201321619
c ESO 2013
Observational evidence for dissociative shocks in the inner 100 AU
of low-mass protostars using Herschel -HIFI
L. E. Kristensen1,2 , E. F. van Dishoeck1,3 , A.O. Benz4 , S. Bruderer3 , R. Visser5 , and S. F. Wampfler6,7
1
2
3
4
5
6
7
Leiden Observatory, Leiden University, PO Box 9513, 2300 RA Leiden, The Netherlands
e-mail: [email protected]
Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA
Max Planck Institut für Extraterrestrische Physik, Giessenbachstrasse 1, 85748 Garching, Germany
Institute for Astronomy, ETH Zurich, 8093 Zurich, Switzerland
Department of Astronomy, University of Michigan, 500 Church Street, Ann Arbor, MI 48109-1042, USA
Centre for Star and Planet Formation, Natural History Museum of Denmark, University of Copenhagen, Øster Voldgade 5-7,
1350 Copenhagen K, Denmark
Niels Bohr Institute, University of Copenhagen, Juliane Maries Vej 30, 2100 Copenhagen Ø, Denmark
Received 1 April 2013 / Accepted 24 June 2013
ABSTRACT
Aims. Herschel-HIFI spectra of H2 O towards low-mass protostars show a distinct velocity component not seen in observations from
the ground of CO or other species. The aim is to characterise this component in terms of excitation conditions and physical origin.
Methods. A velocity component with an offset of ∼10 km s−1 detected in spectra of the H2 O 110 –101 557 GHz transition towards six
low-mass protostars in the “Water in star-forming regions with Herschel” (WISH) programme is also seen in higher-excited H2 O lines.
The emission from this component is quantified and local excitation conditions are inferred using 1D slab models. Data are compared
to observations of hydrides (high-J CO, OH+ , CH+ , C+ , OH) where the same component is uniquely detected.
Results. The velocity component is detected in all six targeted H2 O transitions (Eup ∼ 50–250 K), as well as in CO 16–15 towards one
source, Ser SMM1. Inferred excitation conditions imply that the emission arises in dense (n ∼ 5 × 106 –108 cm−3 ) and hot (T ∼ 750 K)
gas. The H2 O and CO column densities are 1016 and 1018 cm−2 , respectively, implying a low H2 O abundance of ∼10−2 with respect
to CO. The high column densities of ions such as OH+ and CH+ (both 1013 cm−2 ) indicate an origin close to the protostar where the
UV field is strong enough that these species are abundant. The estimated radius of the emitting region is 100 AU. This component
likely arises in dissociative shocks close to the protostar, an interpretation corroborated by a comparison with models of such shocks.
Furthermore, one of the sources, IRAS 4A, shows temporal variability in the offset component over a period of two years which is
expected from shocks in dense media. High-J CO gas detected with Herschel-PACS with T rot ∼ 700 K is identified as arising in the
same component and traces the part of the shock where H2 reforms. Thus, H2 O reveals new dynamical components, even on small
spatial scales in low-mass protostars.
Key words. astrochemistry – stars: formation – ISM: molecules – ISM: jets and outflows
1. Introduction
Star formation is a violent process, even in low-mass protostars
(Lbol < 100 L ). X-rays and UV radiation from the accreting
star-disk system illuminate the inner, dense envelope while the
protostellar jet and wind impinge on the same inner envelope,
causing shocks into dense gas. The implication is that the physical and chemical conditions along the outflow cavities are significantly different from the conditions in the bulk of the collapsing
envelope. Very little is known about the hot (T > 500 K) gas in
low-mass protostars, primarily because the mass of the hot gas
is at most a few % of that of the envelope. Second, few unique,
abundant tracers of the hot gas in the inner envelope exist, save
H2 and CO observed at near-infrared wavelengths (Herczeg et al.
2011), both of which are very difficult to detect towards the
deeply embedded protostars where the AV is 100 (e.g., Maret
et al. 2009).
Herschel is an ESA space observatory with science instruments
provided by European-led Principal Investigator consortia and with important participation from NASA.
The hot gas is most prominently seen in high-J CO observations with the Photodetector Array Camera and Spectrometer
(PACS) on Herschel (Poglitsch et al. 2010; Pilbratt et al. 2010),
where CO emission up to J = 49–48 is detected towards lowmass protostars (Eup ∼ 5000 K; Herczeg et al. 2012; Goicoechea
et al. 2012). The high-J CO emission (Jup > 14) traces two
components with rotational temperatures of 300 and 700–800 K
(a warm and hot component, respectively) seen towards several
tens of low-mass protostars (Green et al. 2013; Karska et al.
2013; Manoj et al. 2013). At present it is unclear whether the
two temperature components correspond to separate physical
components, or whether they are part of a distribution of temperatures, or even just a single temperature and density (Visser
et al. 2012; Neufeld 2012; Manoj et al. 2013; Karska et al.
2013). Moreover, depending on the excitation conditions, the rotational temperature may or may not be identical to the kinetic
gas temperature. Santangelo et al. (2013) and Dionatos et al.
(2013) relate the two rotational-temperature components seen
in CO to two rotational-temperature components seen in H2 rotational diagrams; with its lower critical density (∼103 cm−3
Article published by EDP Sciences
A23, page 1 of 13
A&A 557, A23 (2013)
versus >105 cm−3 for high-J CO transitions) H2 is more likely
to be thermally excited suggesting that the excitation is thermal.
The PACS lines provide little information, beyond the rotational temperature, as they are all velocity-unresolved and no information is therefore available on the kinematics of this gas.
If the very high-J CO emission is caused by shocks in the
inner dense envelope (n 106 cm−3 ) the rotational temperature is similar to the kinetic gas temperature, as proposed by,
e.g., van Kempen et al. (2010) and Visser et al. (2012). Further
support for this hypothesis and the existence of shocks in the
dense inner envelope comes from observations of [O i] at 63 μm
and OH, also done with PACS. Towards the low-mass protostar
HH46, the inferred column densities of these species indicate the
presence of fast ( > 60 km s−1 ) dissociative shocks close to the
protostar (van Kempen et al. 2010; Wampfler et al. 2010, 2013).
As part of the “Water in star-forming regions with Herschel”
programme (WISH1 ; van Dishoeck et al. 2011), Kristensen et al.
(2010, 2012) detected a distinct velocity component towards
six low-mass protostars in the H2 O 110 –101 557 GHz transition
with the Heterodyne Instrument for the Far-Infrared on Herschel
(HIFI; de Graauw et al. 2010). The component is typically blueshifted from the source velocity by ∼2–10 km s−1 and the width
is in the same range as the offset, ∼5–10 km s−1 . Kristensen et al.
(2012) referred to this component as the “medium component”
and associated it with shocks on the inner envelope/cavity wall
based on the coincidence of H2 O masers and this velocity component. The maser association suggests excitation conditions
where the density is >107 cm−3 and T > 500 K (Elitzur 1992),
conditions similar to those inferred for the high-J CO emission.
Yet, so far little is known concerning this velocity component,
its origin in the protostellar system and the local conditions, primarily because the component is not seen in ground-based observations of lower-excited lines towards these same sources.
In this paper, we combine observations of the H2 O offset
component presented above with observations of light hydrides
(OH, OH+ , CH+ ; Wampfler et al. 2013; Benz et al., in prep.)
and highly excited velocity-resolved CO (up to J = 16–15, for
one source) to constrain the physical and chemical conditions in
this velocity component. The component considered in this paper is the same as presented in Kristensen et al. (2012), except
towards one source, NGC 1333-IRAS 2A. The observations and
data reduction are described in Sect. 2. The results are presented
in Sect. 3. In Sect. 4 we discuss the derived excitation conditions
and the interpretation of the velocity component in the context
of irradiated shocks. Finally, Sect. 5 contains the conclusions.
Kristensen et al. (2010, 2012) identified three characteristic
components from the line profiles of H2 O lines in low-mass protostars: a narrow (FWHM < 5 km s−1 ), medium (5 < FWHM <
20 km s−1 ), and broad (FWHM > 20 km s−1 ) component. The
component in this study is characterised by its offset rather than
line width as in our previous work, and thus named “offset” component. For most sources, it is the same as the medium component, except for IRAS 2A. Towards this source the “offset”
component is broader than the “medium” component (40 km s−1
vs. 10 km s−1 ) but is clearly offset from the source velocity.
The offset component is not seen in low-J CO transitions from
the ground but was detected in H2 O emission (Kristensen et al.
2010), which is not the case for the “medium” component, and
hence the redefinition.
2.2. Herschel observations
The central positions of the six low-mass protostars were observed with HIFI on Herschel in twelve different settings cov+
+
ering six H16
2 O and two CO transitions, and HCO , OH, CH ,
OH+ and C+ (Eu /kB ≈ 50−300 K; see Table 1 for an overview).
Only Ser-SMM1 was observed in all settings, the other sources
in a sub-set (Tables 1 and A.1).
Data were obtained using the dual beam-switch mode with
a nod of 3 and a fast chop and continuum optimisation, except for the ground-state ortho-H2O line at 557 GHz, where a
position switch was used (see Kristensen et al. 2012, for details). The diffraction-limited beam size ranges from 12 to 39
(2800–9200 AU for a distance of 235 pc). Data were reduced using HIPE ver. 8. The calibration uncertainty is taken to be 10%
for lines observed in Bands 1, 2, and 5 while it is 30% in Band 4
(Roelfsema et al. 2012). The pointing accuracy is ∼2 . A mainbeam efficiency of 0.65–0.75 is adopted (Table 1). Subsequent
analysis of the data is performed in CLASS2 including subtraction of linear baselines. H- and V-polarisations are co-added after
inspection; no significant differences are found between the two
data sets.
To compare observations done with different beam sizes, all
components are assumed to arise in an unresolved physical component even in the smallest beams (12 ). The emission is scaled
to a common beam-size of 20 (the beam at 1 THz) using a simple geometrical scaling for a point source. We argue a posteriori
that this is an appropriate scaling (Sect. 3.4).
3. Results
2. Observations
3.1. Quantifying emission from the offset component
Six low-mass protostars in NGC 1333 and Serpens clearly show
the presence of an offset component in the H2 O 557 GHz line
(Kristensen et al. 2012). The sources are NGC 1333-IRAS 2A,
IRAS 3A, IRAS 4A, IRAS 4B and Serpens SMM1, SMM3.
Two sources, IRAS 3A and SMM3, show the offset component in absorption against the outflow and/or continuum. The
sources are all part of the WISH sample of low-mass protostars
(van Dishoeck et al. 2011; Kristensen et al. 2012). IRAS 3A
was only observed in the H2 O 557 GHz line and is therefore excluded from further analysis. The component is not seen towards
any other source, and the possible reasons will be discussed in
Sect. 4.3.
Figure 1 shows the H2 O 110 –101 (557 GHz) and 202 –111
(988 GHz) spectra towards all sources with an offset component;
the offset component is marked in the figure. The remaining parts
of each profile will be presented and analysed in a forthcoming
paper (Mottram et al., in prep.). The component is characterised
by being significantly blue-shifted (source − LSR > 2 km s−1 ),
or, for the isolated case of IRAS 4B, by being located close to
the source velocity (source − LSR < 1 km s−1 ). Furthermore, the
FWHM (Δ) is 4 km s−1 extending to 40 km s−1 . These ranges
may be caused by inclination effects, which will be discussed
further in Sect. 4.2. To quantify the emission, each profile is decomposed into Gaussian components, but leaving the deep absorptions seen in the ground-state lines masked out. The resulting offsets and FWHM are shown in Fig. 2. For each source,
1
2
2.1. Source sample
http://www.strw.leidenuniv.nl/WISH
A23, page 2 of 13
http://www.iram.fr/IRAMFR/GILDAS/
L. E. Kristensen et al.: Observational evidence for dissociative shocks in the inner 100 AU of low-mass protostars
Table 1. Species and transitions observed with Herschel-HIFI containing the offset component.
Transition
H2 O
110 –101
212 –101
111 –000
202 –111
211 –202
312 –303
312 –221
CO 10–9
16–15
CH+ 1–0
OH+ 1–0c
C+
2–1
HCO+ 6–5
c
CH
c
OH
ν
(GHz)
556.94
1669.90
1113.34
987.93
752.03
1097.37
1153.13
1151.99
1841.35
835.14
1033.12
1900.54
535.06
536.76
1834.75
λ
(μm)
538.29
179.53
269.27
303.46
398.64
273.19
259.98
260.24
162.81
358.97
290.18
157.74
560.30
558.52
163.40
Eu /kB
(K)
61.0
114.4
53.4
100.8
136.9
249.4
249.4
304.1
751.7
40.1
49.6
91.2
89.9
25.8
269.8
A
(s−1 )
3.46(–3)
5.59(–2)
1.84(–2)
5.84(–3)
7.06(–3)
1.65(–2)
2.63(–3)
1.00(–4)
4.05(–4)
2.3(–3)
1.8(–2)
2.30(–6)
1.27(–2)
6.80(–4)
2.12(–2)
Beama
( )
38.1
12.7
19.0
21.5
28.2
19.7
18.4
18.4
11.5
25.4
20.5
11.2
39.6
39.6
11.6
tint b
(min.)
13.0
23.7
43.5
23.3
18.4
32.5
13.0
13.0
44.6
15.6
30.1
15.8
43.7
43.7
44.6
ηMB a
Sources
0.75
0.71
0.74
0.74
0.75
0.74
0.64
0.64
0.70
0.75
0.74
0.69
0.75
0.75
0.70
IRAS 2A, IRAS 3A, IRAS 4A, IRAS 4B, SMM1, SMM3
IRAS 2A, IRAS 4A, IRAS 4B, SMM1, SMM3
IRAS 2A, IRAS 4A, IRAS 4B, SMM1, SMM3
IRAS 2A, IRAS 4A, IRAS 4B, SMM1, SMM3
IRAS 2A, IRAS 4A, IRAS 4B, SMM1, SMM3
IRAS 2A, IRAS 4A, IRAS 4B, SMM1, SMM3
IRAS 2A, IRAS 4A, IRAS 4B, SMM1, SMM3
IRAS 2A, IRAS 4A, IRAS 4B, SMM1, SMM3
SMM1
IRAS 2A, IRAS 4A, IRAS 4B, SMM1
IRAS 2A, IRAS 4A, IRAS 4B, SMM1
IRAS 2A, IRAS 4A, IRAS 4B, SMM1
IRAS 2A, IRAS 4A, IRAS 4B, SMM1, SMM3
IRAS 2A, IRAS 4A, IRAS 4B, SMM1, SMM3
SMM1
Notes. For lines with hyperfine splitting (OH and OH+ ) only the strongest component is shown here. From the JPL database of molecular spectroscopy (Pickett et al. 1998). (a) Half-power beam width, from Roelfsema et al. (2012). (b) Total on + off integration time incl. overheads.
(c)
Transitions with hyperfine splitting.
Table 2. H2 O emission in the offset component.
IRAS 2A
IRAS 3A
IRAS 4A
IRAS 4B
Ser-SMM1
Ser-SMM3
rmsa
T MB d T peak
T MB d T peak
T MB d T peak
T MB d T peak
T MB d T peak
T MB d T peak
(mK) (K km s−1 ) (K) (K km s−1 ) (K) (K km s−1 ) (K) (K km s−1 ) (K) (K km s−1 ) (K) (K km s−1 ) (K)
110 –101
7
2.55
0.06
–1.25
–0.19
4.23
0.40
3.78
0.89
0.71
0.17
–1.59
–0.11
212 –101
80
7.73
0.17
...
...
2.40
0.23
16.35
3.84
4.51
1.06
–2.08
–0.14
111 –000
16
3.96
0.09
...
...
3.33
0.31
5.78
1.36
1.51
0.36
–1.56
–0.10
202 –111
16
4.65
0.11
...
...
3.67
0.35
2.17
0.51
2.54
0.60
0.69
0.05
211 –202
18
2.42
0.06
...
...
1.96
0.18
1.44
0.34
1.37
0.32
0.52
0.04
312 –303
56
3.14
0.07
...
...
1.17
0.11
1.11
0.26
1.66
0.39
0.25
0.02
312 –221
70
3.16
0.07
...
...
2.09
0.20
1.16
0.27
2.66
0.63
<0.05
<0.01
Δ (km s−1 )
40
6
10
4
4
14
–5
5
–1
8
4
2
LSR (km s−1 )
–12.7
–3.3
–8.0
0.9
–4.5
–5.6
offset b (km s−1 )
Transition
Notes. Obtained from Gaussian fits to each component; negative values are for absorption. Upper limits are 1σ. (a) Measured in 1 km s−1 channels.
(b)
Offset velocity with respect to the source velocity as reported by Yıldız et al. (2013).
a characteristic offset velocity and FWHM were chosen based
on the decomposition of high-S/N data without self-absorption,
typically the 312 –303 (1097 GHz) and 202 –111 (988 GHz) transitions (Fig. 2). These parameters were fixed and the decomposition redone for all spectra, letting only the intensity and the
parameters of the secondary component be free. The secondary
component is typically the broader component (except for the
case of IRAS2A) associated with the outflow (Kristensen et al.
2010, 2012). The resulting intensities are listed in Table 2.
The main uncertainty in the listed intensities comes from the
fitting and the uniqueness of the fit, particularly for low signalto-noise lines. For strong components and very offset components, the fit is unique and the corresponding uncertainty low, as
illustrated by the scatter of the width and offset shown in Fig. 2.
By fixing the width and offset we have removed this scatter by
assuming that the shape of the component is independent of excitation. The decompositions obtained by fixing the width and
offset are equally good to the decompositions where all parameters are free, where the quality of the fit is taken to be the residual. Typically, the rms of the residual is <1.5 times the rms in a
line-free region.
CO 10–9 spectra were also examined for the presence of an
offset component by using the same kinematic parameters as in
the H2 O decomposition. The strongest CO 10–9 line is observed
towards SMM1 (Yıldız et al. 2013) where no direct evidence
is found for an offset component; the line profile can be decomposed without the need for an additional offset component.
SMM1 does show a clear offset component in CO 16–15, and
the question therefore arises, how much emission from the offset component can be hidden in the CO 10–9 profile? Figure 4
shows the CO 16–15, 10–9 and H2 O 312 –303 spectra obtained
towards SMM1. The offset component was first fitted using
CO 16–15 and subsequently the Δ and LSR were fixed and used
to quantify emission in the CO 10–9 profile. Second, the same
exercise was done but the offset parameters from the H2 O decomposition were fixed. There is little difference in terms of the
integrated intensity no matter whether the CO or H2 O parameters are chosen (5.5 vs. 4.9 K km s−1 ). The quality of the fits
are the same as if the offset component is not included, and we
therefore treat the inferred intensities as upper limits. Similarly
for the other sources, only upper limits are available and these
are based on the H2 O kinematic parameters (Table 3).
A23, page 3 of 13
A&A 557, A23 (2013)
5
IRAS2A
IRAS3A
IRAS4A
IRAS4B
SMM1
SMM3
υoffset (km s−1)
0
−5
−10
−15
−20
0
10
20
30
FW HM (km s−1)
40
50
Fig. 2. Velocity of the offset component with respect to the source velocity as a function of FWHM for all observed transitions. The plus
signs show the results of the decomposition of each line, whereas the
circles show the values of the offset and width chosen as fixed.
Table 3. CO 16–15 and limits on CO 10–9 integrated intensities in the
offset components.
Source
IRAS 2A
IRAS 4A
IRAS 4B
SMM1(H2 O)a
SMM1(CO)b
SMM3
CO 10–9
T MB d
(K km s−1 )
<0.9
<1.2
<6.6
<4.9
<5.5
<2.3
T MB d
(K km s−1 )
...
...
...
...
6.2
...
CO 16–15
LSR
(km s−1 )
5.5
Δ
(km s−1 )
3.4
Notes. Upper limits are 1σ and are obtained by fixing the position and
width of the offset component from the H2 O data. (a) LSR and FWHM
from the H2 O profiles. (b) LSR and FWHM from the CO 16–15 profile.
Fig. 1. Top: continuum-subtracted HIFI spectra of the H2 O 110 –101
ground-state transition at 557 GHz (Eup = 60 K). The profiles have been
decomposed into Gaussian components and the best-fit profile is shown
in blue. The offset component is highlighted in magenta for clarity. The
baseline is shown in green and the source velocity with a red dashed
line. Bottom: same as the top figure, except that the line shown is the
excited H2 O 202 –111 line at 988 GHz (Eup = 100 K). More details on
the decomposition are found in Kristensen et al. (2010, 2012) for the
NGC1333 sources and Ser SMM3.
CH+ and OH+ spectra show the offset component in absorption (Benz et al., in prep.). OH+ is detected towards all sources
whereas CH+ is detected towards two out of four sources.
IRAS 4B shows the CH+ profile to be shifted both with respect to the source velocity and the offset component seen
A23, page 4 of 13
in H2 O; it is centred on 5.7–6.0 km s−1 and is thus offset by more
than 1 km s−1 towards the blue. The CH+ feature towards SMM1
is centred on the velocity of the offset component and has a larger
FWHM than the offset component seen in H2 O and CO (Fig. 3).
C+ is detected towards IRAS 2A and SMM1, in both cases blueshifted with respect to the source velocity by ∼4–5 km s−1 .
OH is detected towards Ser SMM1 with HIFI at 1835 GHz
(Wampfler et al., in prep.). A Gaussian decomposition is complicated by the fact that the hyperfine transitions are very closely
spaced (2.4 km s−1 ). Nevertheless, by fixing the intensity ratios
of the hyperfine components, i.e., assuming the emission is optically thin and that the intensity ratios scale with Aul gu , a fit is
obtained and the OH emission likely contains a mixture of the
offset and broad component. Because of the shape of the profile, further decomposition is not performed, instead we assume
that 50% of the emission can be attributed to the offset component. HIFI OH spectra are not available towards the other sources
as part of WISH.
3.2. Time variability
IRAS 4A was re-observed in the H2 O 312 –303 transition at 1097
GHz as part of an OT2 programme (PI: Visser) on Aug. 2, 2012,
nearly two years after the original observations (July 31, 2010).
L. E. Kristensen et al.: Observational evidence for dissociative shocks in the inner 100 AU of low-mass protostars
Fig. 3. Continuum-subtracted HIFI spectra of H2 O, CO, OH, CH+ ,
OH+ obtained towards the central position of Ser SMM1 (source =
8.5 km s−1 ). The red vertical line indicates source , while the blue dashed
line shows the position of the offset component. The blended OH triplet
is centred on the strongest hyperfine component as indicated by the
three vertical black lines situated directly beneath the OH spectrum.
The same is the case for the OH+ spectrum, with the location of the hyperfine components and their relative strengths indicated by the black
lines above the spectrum. The spectra have been shifted vertically for
clarity and in some cases scaled by a factor, indicated on the figure.
CH+ and OH+ are fitted by a single Gaussian, OH by two, and CO and
H2 O by three; the offset components are shown in magenta.
During that period, the offset component doubled in intensity
(Fig. 5). IRAS 2A, IRAS 4B and SMM1 were also re-observed
as part of the same OT2 programme, but show no signs of variability. All H2 O observations towards IRAS 4A presented here
were performed over a period of two months, and we assume
that no significant variability took place over that time period.
To verify that the change in emission is not caused by a pointing offset when the data were obtained, the pointing offsets were
checked using HIPE 9.1. The recorded pointing offset towards
IRAS 4A was 3. 8 in July 2010 and less than 0. 5 in 2012. The
pointing offsets were similar towards IRAS 4B for both epochs
and 2. 5 for both epochs for the other sources. For the pointing
offset to be the cause of the intensity difference, the offset component would need to be located at the edge of the HIFI beam,
i.e., at a distance of more than 10 from the pointing centre to
cause a doubling in intensity. Otherwise the pointing alone cannot account for the change. Below in Sect. 4.1 we argue why this
origin is unlikely, based on the hydride absorption.
HIFI has an inherent calibration uncertainty of ∼10%
(Roelfsema et al. 2012). However, only the offset component
shows a noticeable difference in intensity, the broader underlying
outflow component appears unchanged between the two epochs.
Towards SMM1, on the other hand, both the broad and offset
components change intensity slightly, a change which can be attributed to calibration uncertainties; the difference in intensity
Fig. 4. Decomposition of the CO 10–9 profile towards SMM1 using either the best-fit parameters obtained from CO 16–15 (top) or H2 O (bottom). The Gaussian profiles illustrate the maximum amount of emission
that can be hidden in the CO 10–9 profile.
is 10% across the spectrum. In conclusion, the most likely explanation is that the offset component seen towards IRAS 4A
changed in intensity over the past two years.
Spectra of species other than H2 O towards IRAS 4A were
obtained over a period of 6 months from March 3, 2010 to
September 3, 2010, and we assume that little or no change took
place in that time frame.
3.3. CO and limits on kinetic temperature
The offset component is seen in CO J = 16–15 towards SMM1,
but not in the lower-excited J = 10–9 line. Thus, the upper
limit from J = 10–9 can be used to provide a lower limit on
the rotational temperature and a corresponding upper limit on
the column density, assuming the level populations are in local thermodynamic equilibrium and optically thin. The results
are illustrated in Fig. 6. The upper limit on the column density
is ∼7 × 1013 cm−2 in the 11. 5 beam of the CO 16–15 transition.
Assuming LTE and that both lines are optically thin, the limit
on the rotational temperature is T rot 270 K. Goicoechea et al.
(2012) find that the CO ladder towards Ser SMM1 from J = 4–3
to 49–48 consists of three rotational-temperature components
with T rot = 100, 350 and 600 K, respectively, corresponding to
low-J, mid-J and high-J CO emission (J 14, 26 and 42, respectively). The offset component is clearly not associated with
the 100-K temperature component, but based on the limits on the
rotational temperature it is not possible to conclude whether it is
A23, page 5 of 13
A&A 557, A23 (2013)
32
260 K
ln(Nup/gup)
31
280 K
J = 10−9
30
29
28
0
CO 10-9 limit from:
CO 16−15
H2 O
200
400
Eup/kB (K)
J = 16−15
600
800
Fig. 6. CO rotational diagram for Ser SMM1 based on the upper limit of
CO J = 10–9 emission obtained from two different decompositions, and
the detection in CO J = 16–15. The lower limits on the rotational temperatures are shown for each upper limit on the CO J = 10–9 emission.
Fig. 5. H2 O 312 –303 spectra at 1097 GHz towards SMM1, IRAS 2A,
IRAS 4A and IRAS 4B observed at two different epochs. The offset
component towards IRAS 4A has doubled in intensity, as shown by the
two Gaussian fits in magenta. All spectra are shifted such that the source
velocity is at 0 km s−1 .
associated with the warm or hot component. Indeed, if the distribution is continuous it is possible that the rotational temperatures
do not correspond to discrete temperature regimes (e.g., Neufeld
2012).
For the specific case of SMM1 it is clear that the offset component is a distinct physical component based on the line profile,
and this component must be present in the observed CO ladder.
Moreover, the contribution is embedded in the integrated emission for Jup > 10 which further illustrates the need for high spectral resolution observations to isolate emission from the separate
dynamical components, as opposed to observations with, e.g.,
SPIRE and PACS on Herschel. The same is likely the case for
the other sources although a future analysis will show to what
extent the CO J = 10–9 data can be used to constrain the CO rotational temperature. If emission is strongly beam-diluted (see
next section) it is likely that high angular resolution observations
with a facility such as ALMA using CO J = 6–5 will be able to
further constrain the rotational temperature as well.
3.4. H2 O and CO excitation conditions
To determine the H2 O excitation conditions, n(H2 ), T
and N(H2 O), specific H2 O line ratios are examined. The
H2 O 312 –303 /312 –221 ratio is particularly useful in providing initial constraints on the column density. Because the two transitions share the same upper level, the ratio is straightforward to
calculate in the optically thin limit and is equal to 6.7. However,
A23, page 6 of 13
observations of these two transitions at 1153 and 1097 GHz, i.e.,
in similar beams, reveal an intensity ratio in the offset component ranging from 0.6 (IRAS 4A) to 1.7 (SMM3). Such a ratio
can only be explained if the 312 –303 transition is optically thick
(τ > a few) and the column density is greater than ∼1016 cm−2
for any given emitting area.
Figure 7 shows the H2 O 202 –111 /211 –202 versus H2 O 312 –
303 /312 –221 line ratios for various H2 O column densities (4 ×
1015 –1017 cm−3 ), H2 densities (106 –109 cm−3 ) and a temperature of 750 K. The ratios are calculated using the non-LTE statistical equilibrium code RADEX (van der Tak et al. 2007) for
line widths of Δ = 4, 10, 14 or 40 km s−1 corresponding to the
FWHM of the different offset components. The H2 O-H2 collisional rate coefficients from Daniel et al. (2011) are used and the
H2 and H2 O o/p ratios are set to 3, the high-temperature equilibrium value. Observed line ratios are also shown and these have
been scaled to the same 20 beam assuming that the emitting
region is much smaller than the beam.
The resulting line ratios typically change by less than 10%
for temperatures in the range of 500–1000 K, i.e. they only
weakly depend on the assumed temperature. Goicoechea et al.
(2012) find from an excitation analysis of more H2 O lines that
a kinetic temperature of ∼800 K reproduces the H2 O line ratios
as well as the high-J part of the CO ladder observed towards
SMM1. We choose to fix the temperature to 750 K, the halfway
point between 500 and 1000 K but note that this value is not
constrained by the H2 O data.
For most model results there is a degeneracy between a (relatively) low H2 density (∼5 ×106 cm−3 ), low H2 O column density
(a few times 1016 cm−2 ) and a high H2 density (>108 cm−3 ), high
H2 O column density (>1017 cm−2 ) (Fig. 7). This degeneracy is
most evident when comparing model results to the observations
of IRAS 4A, IRAS 4B and SMM1, and corresponds to whether
the line emission is sub-thermally or thermally excited. For the
high column density case, the ground-state H2 O lines are very
optically thick, τ > 100, which may affect the accuracy of the
radiative transfer. In the following, the results with the lowest
column density and thereby lowest opacity will be analysed. The
model results are summarised in Table 4.
The offset component was linked with 22 GHz H2 O maser
emission in Kristensen et al. (2012). For H2 O to mase, a density
L. E. Kristensen et al.: Observational evidence for dissociative shocks in the inner 100 AU of low-mass protostars
I4B, SMM1:
Δυ = 4 km s−1
1.75
I4A:
Δυ = 10 km s−1
1.50
1.25
SMM1
1.00
I4B
0.75
0.50
106
107
108
109 cm−3
H2O 202 − 111 / 211 − 202
1.75
1.50
SMM3:
Δυ = 14 km s−1
I2A:
4.0×1015
Δυ = 40 km s−1
1.3×1016
4.0×1016
1.3×1017 cm−2
1.25
1.00
0.75
0.50
0
1
2
H2O 312 − 303 / 312 − 221
3 0
1
Table 4. H2 O and CO excitation conditions.
Source
IRAS 2A
IRAS 4A
IRAS 4B
SMM1
SMM3
n(H2 )
(cm−3 )
5 × 106
5 × 106
1 × 107
5 × 106
5 × 107
N(H2 O)a
(cm−2 )
4 × 1016
1 × 1016
4 × 1015
4 × 1016
1 × 1016
N(CO)a
(cm−2 )
1 × 1018
T kin b
(K)
750
750
750
750
750
r
(AU)
80
140
160
110
50
Notes. (a) Column density over the emitting region with radius r.
(b)
Kinetic temperature fixed in the model.
of ∼107 cm−3 is required (Elitzur 1992), which is typically a factor of two higher than what is inferred here. The maser density
is based on H2 O-H2 collisional rate coefficients which are more
than twenty years old. Furthermore, the error bar on the observed
line ratios is such that the best-fit densities span an order of magnitude and a density of 107 cm−3 cannot be excluded. Thus the
conclusion that the offset component is coincident with masers
remains unchanged.
The absolute H2 O 202 –111 intensity from the RADEX models are compared to the observed intensity in the offset component to estimate the beam filling factor, or, alternatively, the
radius of the emitting region, r. Finally, the CO column density
is varied until the CO 16–15 intensity towards SMM1 is recovered for the same conditions as for H2 O. Typically, the radius
of the emitting region is of the order of 100 AU (Table 4), or
about ∼0. 5 at a distance of 235 pc.
Towards SMM1, Goicoechea et al. (2012) obtain CO column
densities of the warm (T ∼ 375 K) and hot (∼800 K) components of 1018 cm−2 and 5 × 1016 cm−2 , respectively, over a region
with a radius of 500 AU. Note that these temperatures are inferred from modelling the CO ladder and H2 O emission, and are
not identical to the measured CO rotational temperature (600 K,
Goicoechea et al. 2012). If our inferred CO column density
of 1018 cm−2 over a 110 AU emitting radius is scaled to a radius
2
3
Fig. 7. H2 O 202 –111 / 211 –202 versus H2 O
312 –303 / 312 –221 line ratios for various H2 O column densities and H2 densities from RADEX
models. The four panels are for different line
widths corresponding to the width of each of
the offset components. The temperature is fixed
at 750 K. The different line styles correspond
to different H2 O column densities and the dots
are for different H2 densities. The observed ratios are scaled to the same beam and marked
with 15% error bars in both ratios.
of 500 AU, the CO column density becomes ∼5 × 1016 cm−2 .
It is therefore likely that the hot component observed in the
high-J CO data is identical to the offset component identified
in the HIFI H2 O and CO data; the column density of the warm
CO component is too high to be hidden in the HIFI offset component. This analysis shows that the profiles are necessary for
disentangling the different kinematical components in each spectrum, and that the hot CO detected with PACS is likely a distinct
physical component towards this source.
This work assumes that the temperature of the H2 O emitting gas is the same as that of CO, and that the temperature of the H2 O emitting gas is close to what has been determined by, e.g., Goicoechea et al. (2012). What if this is not
the case? Furthermore, how much column density can be hidden in the CO 10–9 profiles, where the offset component is
not detected? The CO J = 10–9 contribution of the warm
and hot components is estimated from the rotational diagrams
in Karska et al. (2013), Herczeg et al. (2012) and Goicoechea
et al. (2012). In all cases, the integrated CO J = 10–9 intensity
is of the order of ∼30–40 K km s−1 for the warm component
and ∼2.5–3 K km s−1 for the hot component when extrapolating
the linear fits from J ∼ 15–40 down to J = 10 and assuming
emission is optically thin. The opacity can be estimated from
optically thin 13 CO 10–9 emission and the J = 10–9 line is optically thin away from the line centre (San José-García et al. 2013;
Yıldız et al. 2013). The upper limit for the offset component in
the CO 10–9 data is ∼1–6 K km s−1 (Table 3). Thus, the offset
component would have easily been detected in the HIFI spectra
if it were associated with the warm PACS component, but not if
it is associated with the hot component.
The H2 O/CO abundance ratio towards SMM1 is ∼0.04. Our
analysis assumes that CO and H2 O share excitation conditions.
This may not be the case, as CO is shifted with respect to H2 O
by ∼1.5 km s−1 towards SMM1, although they have identical line
widths. Goicoechea et al. (2012) find a higher H2 O/CO abundance ratio of 0.4 for the same H2 density of 5 × 106 cm−3 ,
but for a different emitting region. Since little H2 data exist towards the central source position, and certainly no H2 data with
A23, page 7 of 13
A&A 557, A23 (2013)
Table 5. Column densities of OH+ , C+ and CH+ where detected, and 3σ upper limits on CH and HCO+ column densities.
Source
IRAS 2A
IRAS 4A
IRAS 4B
SMM1
SMM3
N(OH+ )
(cm−2 )
2.2 × 1013
1.5 × 1013
7.5 × 1012
3.3 × 1013
...
N(CH+ )
(cm−2 )
<4.2 × 1011
<6.4 × 1011
1.8 × 1012
2.4 × 1013
...
N(C+ )
(cm−2 )
3.9 × 1017
<9.6 × 1016
<3.9 × 1017
4.2 × 1016
...
N(HCO+ )
(cm−2 )
<8.0 × 1013
<6.6 × 1013
<2.0 × 1014
<3.3 × 1014
<1.1 × 1015
N(CH)
(cm−2 )
<3.7 × 1014
<3.1 × 1014
<9.4 × 1014
<1.5 × 1015
<4.9 × 1015
Notes. OH+ and CH+ are from Benz et al. (in prep.). SMM3 was not targeted for observations of OH+ , CH+ and C+ .
the velocity resolution required to isolate the offset component,
the H2 O/CO ratio serves as a proxy for the H2 O abundance
with respect to H2 . For a canonical CO abundance of 10−4 the
H2 O abundance is 4 × 10−6 and thus lower by several orders of
magnitude compared to what would be expected if all oxygen
were locked up in H2 O (∼3 × 10−4 ).
3.5. OH+ , C+ and CH+
Observations and characteristics of the hydride observations are
reported in Benz et al. (in prep.). We here summarise the main
results concerning the offset component and report again the hydride column densities for completeness.
The offset component is uniquely identified in absorption in
both OH+ and CH+ towards IRAS 4B and SMM1, and in OH+
towards all sources. The velocity offsets are consistent with those
seen in H2 O and, for the case of SMM1, in CO J = 16–15. From
the absorption features it is possible to directly measure the absorbing column through
8π ν3 glow
τ d ,
(1)
Nlow = 3
c Aul gup
where ν is the line frequency, Aul the Einstein A-coefficient and g
the statistical weight of the lower and upper levels. The opacity,
τ, is determined as τ = ln (T cont /T line ). In determining the column
density, it is implicitly assumed there is no re-emission of the absorbed photons. The measured values are given in Table 5 along
with 3σ upper limits (Benz et al., in prep.).
The CH+ column densities are in the range of < a few
times 1011 cm−2 to 2 × 1013 cm−2 . This range is similar to that
found in diffuse interstellar clouds (∼1012 –1014 cm−2 ; Gredel
1997). The OH+ column densities are of the order of 1013 cm−2 .
The OH+ column density is similar to what Bruderer et al.
(2010b) measured towards the high-mass star-forming region
AFGL2591, N(OH+ ) ∼ 1.6 × 1013 cm−2 , whereas the CH+ column densities towards the low-mass objects are 1–2 orders of
magnitude lower than towards AFGL2591, N(CH+ ) ∼ 1.8 ×
1014 cm−2 .
C+ is not uniquely identified with the offset component although an absorption feature is seen towards SMM1 at the velocity of the offset component, and towards IRAS 2A closer
to the source velocity. The measured column densities are 4 ×
1017 cm−2 . The observations were performed in dual-beamswitch mode and it cannot be ruled out that some emission
is missing because it is chopped out. For example, Larsson
et al. (2002) observed extended [C ii] emission over the entire
Serpens core with ISO-LWS. However, the absorption is expected to mainly affect any kinematical component at or close to
the source velocity; the offset component is blue-shifted by several km s−1 which suggests that the [C ii] absorption is unrelated
to the cloud or Galactic foreground but intrinsic to the sources.
A23, page 8 of 13
3.6. OH, CH and HCO+
The offset component is not uniquely identified in the spectra
of OH, CH and the very deep HCO+ J = 6–5 line obtained
serendipitously in our observations of H18
2 O 110 –101 (Kristensen
et al. 2010). Nevertheless, for the OH HIFI spectrum towards
SMM1, we assume that 50% of the emission and therefore 50%
of the column density can be assigned to the offset component (see discussion above, Sect. 3.1). An OH column density
of 5 × 1015 cm−2 is adopted, a value which is probably accurate
to within a factor of a few (Wampfler et al. 2013).
As for CH and HCO+ , emitting region sizes, temperatures,
H2 densities and line-widths are as for H2 O and CO. For the case
of CH, only one transition is observed and we therefore adopt the
upper limit on the rotational temperature from Bruderer et al.
(2010b) of 25 K to estimate the total column density. If CH
is in LTE and is optically thin with a rotational temperature
of >500 K, the upper limit is only a factor of ∼5 higher than
what is given in Table 5. The same procedure is adopted for
HCO+ 6–5 with the exception that the rotational temperature
is taken to be >500 K, because the critical density of HCO+
is ∼2 × 107 cm−3 and so emission may be in LTE. If HCO+ is
strongly subthermally excited and the rotational temperature is
only 25 K, the column densities are overestimated by ∼25%.
Typical upper limits are ∼1014 cm−2 and 1015 cm−2 for HCO+
and CH, respectively.
4. Discussion
4.1. Location of the offset component
Because the offset component is typically blue-shifted and because it sometimes appears in absorption in certain species
against the continuum, it is possible to constrain the physical location of the component in the protostellar system. First, we assume that the offset component consists of both a red- and blueshifted component, but that the red-shifted component is hidden
from view by some obscuring agent. This obscuring agent may
consist of either gas or dust (or both), and both possibilities are
discussed below.
The CO 16–15 emission is optically thin (Goicoechea et al.
2012) and the red-shifted counterpart is the most difficult to hide.
Is it possible to have a layer of CO gas between the blue- and
red-shifted offset components that could shield the red-shifted
component, and if so, what would the conditions need to be?
For high densities of 5 × 106 cm−3 , the optical depth of the
CO 16–15 line is nearly independent of density, and thus only
depends on temperature and column density. For temperatures
greater than ∼100 K, the optical depth scales almost linearly with
temperature. Thus, for N(CO) = 3 × 1018 cm−2 and T = 750 K
the CO 16–15 transition is optically thick with τ = 3, for
Δ = 4 km s−1 , the width appropriate for SMM1. However, such
L. E. Kristensen et al.: Observational evidence for dissociative shocks in the inner 100 AU of low-mass protostars
104
d(τdust = 1)
103
d (AU)
conditions yield significant emission in the lower-J transitions,
and although the component would remain hidden in higher-J
lines, it would appear in lower-J lines. Only for low temperatures and high column densities, e.g., N(CO) = 3 × 1019 cm−2
and T = 100 K is emission in both the high-J and low-J
lines optically thick, similar to the conditions observed in a
Herbig Ae/Be disk (Bruderer et al. 2012) where even CO 16-15
is marginally optically thick. For H2 densities of 5 × 106 cm−3 ,
a column length of 400 AU is required to obscure this component from view. If the density of the gas is higher, 108 cm−3 , the
column length is correspondingly lower, 100 AU. Such length
scales correspond to the inferred sizes of embedded disks around
Class 0 objects (e.g., Jørgensen et al. 2007, 2009) and is also
similar to the inferred size of the emitting region (Table 4). Gas
densities in excess of 108 cm−3 are only found close to the protostar in the disk, where the temperature is low. Thus it is possible
that the red-shifted counterpart is hidden by large amounts of
high-density, cold CO gas.
The red counterpart of the offset component is also obscured
in H2 O emission. However, for similar conditions as discussed
above (n = 5 × 106 cm−3 , T = 100 K) an H2 O column of
only ∼1016 cm−2 is required to obscure material. If the gas and
dust temperatures are equal and close to 100 K, the water abundance is expected to be within a factor of a few of the CO abundance, and thus it is straightforward to hide any water emission
appearing on the red side of the spectrum by colder H2 O gas.
The shielding molecular gas can be associated with the entrained gas in the molecular outflow. If that is the case, then a significant fraction of the outflow is entrained on very small scales
(<a few hundred AU) in order to hide the red offset component
which also resides in the inner few hundred AU. Alternatively,
the obscuring gas originates in the infalling envelope. Freefalling gas towards a protostar with a mass of 0.5 M has a velocity of 3 km s−1 at a distance of 100 AU, and therefore it is
possible that the gas flowing towards the protostar (red-shifted
and located on the side facing us) shields the red offset component located on the far side of the system in the sources where
the offset and width is low. For the case of IRAS 2A where the
offset is 13 km s−1 and the width is 40 km s−1 , the infalling gas
cannot shield the red-shifted offset component.
Shielding by the dust is another possibility. The lowest frequency at which the offset component is detected is
at 557 GHz in the H2 O 110 –101 transition. Adopting a dust opacity of 5 cm2 g−1 , the opacity from Table 1, Col. 5 of Ossenkopf
& Henning (1994) at 500 μm, a gas/dust ratio of 100 and
a mean molecular weight of 2.8 mH , an H2 column density
of >1024 cm−2 is required for a dust optical depth of 1. At
shorter wavelengths, the dust opacities increase and thus a lower
dust column density is required to shield the offset component.
Figure 8 illustrates the dust τ = 1 surface as a function of wavelength from the inside of the envelope of SMM1 using the spherical envelope model of Kristensen et al. (2012). In this representation an observer is able to see the other side of the envelope if
he is located in the “optically thin” zone; in the “optically thick”
zone the dust blocks emission from the other side. At 162 μm,
the wavelength of the CO J = 16–15 transition, an H2 column
density of more than 1023 cm−2 is required for the dust to be
optically thick, which is obtained on scales of ∼220 AU in this
model, i.e., comparable to the size of the emitting region (Fig. 8).
Therefore it is possible that dust obscures the red-shifted offset
component at 162 μm, but this is not the case at 500 μm.
Jørgensen et al. (2005) showed that on scales of a few
hundred AU, an additional component is required to reproduce interferometric observations. This component is likely the
demit
102
Optically thick
λ = 162 μ m
(CO 16–15)
Optically thin
101 1
10
102
λ (μ m)
103
Fig. 8. Thickness of the optically thick dust zone as a function of wavelength from the inside of the envelope of SMM1 towards the outside
using the model envelope of Kristensen et al. (2012). The diameter of
the emitting region is marked.
protostellar disk, and the column density is sufficient to provide
enough shielding from the far side (see Enoch et al. 2009, for the
example of SMM1). In conclusion, the red-shifted component is
likely hidden either by high-density molecular gas in the disk or
the inner, dense envelope.
The fact that several species tracing the offset component
(e.g., OH+ , CH+ and H2 O towards IRAS 3A and SMM3)
only appear in absorption against the continuum from the disk/
envelope places the offset component in front of the disk. The
continuum-emitting region at the wavelengths where these absorptions appear (300 μm for OH+ to 540 μm for H2 O) is
typically up to ∼500 AU in size (∼2 ; Jørgensen et al. 2009;
Goicoechea et al. 2012). These considerations all point to a physical origin of the offset to within the inner few hundred AU of
the protostar, and located between us and the protostar itself.
4.2. Inclination
The three NGC 1333 sources have well-constrained inclination
angles with respect to the plane of the sky. The outflow from
IRAS 2A is close to the plane (i ∼ 90◦ ), IRAS 4A is at i ∼ 45◦ ,
and IRAS 4B is seen nearly pole on (i ∼ 0◦ ). Fitting the spectral energy distribution of SMM1, Enoch et al. (2009) find that
SMM1 has an inclination of 30◦ . IRAS 2A has the largest offset
and FWHM whereas that towards IRAS 4B shows the smallest
offset and FWHM, with IRAS 4A and SMM1 falling between
these two extremes. This four-point correlation suggests that the
offset and width depend on inclination, and that the offset component is moving nearly perpendicular to the large-scale outflow
as explained below.
If the origin of the offset component is a shock as suggested
by the width and offset of the profile, how will the shape of the
profile depend on the inclination? The offset velocity follows
the inclination as sin(i); when the inclination is 0◦ (the case of
IRAS 4B) the offset is 0 km s−1 whereas it reaches its maximum value at i = 90◦ (IRAS 2A). The width will be narrow
when observing the shock from an orientation close to face-on
(IRAS 4B) because only the velocity component inherent to the
shock is probed. For an edge-on orientation the profile will appear broader because the shock is now observed at inclinations
ranging from the plane of the sky to the line of sight.
A23, page 9 of 13
A&A 557, A23 (2013)
1018
i = 0◦
1017
10
1016
0.2
0.0
−30 −25 −20 −15 −10 −5
υ (km s−1)
1013
υ = 80 km s−1
1012
n = 106 cm−3
1011
0
5
10
Fig. 9. Toy model of line profiles originating in an expanding half
annulus as a function of inclination to the observer. An inclination
of 0◦ corresponds to the annulus expanding into the plane of the sky
(IRAS 4B) and 90◦ corresponds to expansion along the line of sight
(IRAS 2A). The incremental inclination is 10◦ . The inset shows the
measured FWHM of the different profiles as a function of inclination,
i.e., not obtained from a Gauss fit. The annulus expands at 15 km s−1
and the internal velocity dispersion (FWHM) is 9 km s−1 .
To illustrate how the velocity offset and width changes with
inclination in the proposed scenario, we make a geometrical toy
model. The model consists of a half annulus which expands from
the plane of the sky towards the observer. The intensity from any
given point along the annulus corresponds to a Gaussian with
a predefined offset (expansion velocity) and width (internal velocity dispersion) which both stay constant along the annulus.
The offset, however, moves and an expansion into the plane of
the sky corresponds to an offset velocity of 0 km s−1 , i.e., the
offset velocity along the annulus scales with the angle θ going
from 90◦ (the line of sight) to 0 and 180◦ (the plane of the sky).
The annulus is furthermore given an inclination to the plane of
the sky, i, which ranges from 0◦ (the plane of the sky) to 90◦ (the
line of sight). The resulting profiles from the expanding half annuli are shown in Fig. 9, where also the evolution of the width
with inclination is shown. No single set of parameters (offset
and width) are able to reproduce all observations, either because
no such single set exists, or the toy model is too simplistic to
capture the geometry and excitation. For example, here only a
single shock or expansion velocity is considered; a range of velocities may be more appropriate as in the case of bow shocks
(e.g. Kristensen et al. 2007). Nevertheless, the behaviour of the
offset components is captured qualitatively which suggests that
the wind scenario is a geometrically possible solution, and that
a shock velocity of 15 km s−1 is reasonable from the line profile
perspective.
4.3. Physical origin
Both the offset (2–15 km s−1 ) and the width (∼4–40 km s−1 )
are indicative of a shock origin, a shock appearing close to
the protostar. Neufeld & Dalgarno (1989) modelled fast, dissociative shocks and included the effects of UV radiation generated in the shock itself through Lyα emission. After initial
heating to >5 × 104 K, the compressed, shocked gas cools to
T ∼ 5000 K where it reaches a plateau. During this phase, the
electron abundance is high (10−2 ) and molecular ions are abundant, e.g., OH+ and CH+ . The gas is compressed by a factor
A23, page 10 of 13
1014
C+
i = 90◦
1015
HCO+
90
OH+
60
CH+
0.4
i (◦ )
CH
30
OH
0
H2 O
5
CO
0.6
15
N (cm−2)
Intensity (arb. units)
0.8
FW HM (km s−1 )
1.0
Fig. 10. Comparison of inferred column densities over the size of the
emitting region and upper limits with shock model results from Neufeld
& Dalgarno (1989) for a pre-shock density of 106 cm−3 and shock velocity of 80 km s−1 . Observations are marked with black dots and arrows are for upper limits. For the cases where both a black dot and arrow are present (OH+ and C+ ) the dot marks the detection and the arrow
the upper limit towards the other sources; when two arrows are present
(CH and HCO+ ) they illustrate the range of upper limits. Model results
are shown as red circles and are normalised to the inferred CO column
density.
of ∼400 in this stage to ∼108 cm−3 . Eventually the OH formation rate exceeds the destruction rate, and OH brings the temperature down to ∼500 K, at which point the temperature reaches
another plateau while H2 forms. Once H2 is formed, the temperature quickly drops to ∼100 K when CO and H2 O take over as
dominant coolants.
The predicted column densities (Neufeld & Dalgarno 1989)
are in good agreement with the inferred observational column
densities (Fig. 10) for a dense, dissociative shock (106 cm−3 )
with a velocity of 80 km s−1 . There is a trend in the model predictions for higher column densities of C+ with higher velocity, but column densities are not reported for > 60 km s−1 and
n = 106 cm−3 (C+ and CH column densities are taken from the
model with = 60 km s−1 ). In general, the agreement is remarkable and shows that a fast, dense shock is a possible explanation
for the observed column densities.
The high density required by the model is easily found
in the inner parts of the molecular envelope. The high velocity required is probably attained in either the jet or the strong
wind from the protostar, although no direct observations exist of the wind in Class 0 objects. In the following, the “jet”
refers to the highly collimated and fast component observed as
extremely high-velocity features in molecular species, and the
“wind” refers to the wide-angle, slower component seen towards
Class I and II sources (e.g., Arce et al. 2007, and references
therein). The slower wind is primarily observed in forbidden
atomic and ionic transitions at near-infrared and shorter wavelengths (Arce et al. 2007; Ray et al. 2007) where the velocity
is 10–20 km s−1 .
If the shock is moving at 80 km s−1 perpendicular to the outflow direction (see above, Sect. 4.2) the envelope will quickly
dissipate on timescales of <103 years (e.g., Shang et al. 2006)
and the wind would be much faster than what is observed at
later evolutionary stages. The key ingredients that a successful
model should reproduce are: (i) the excitation conditions and
L. E. Kristensen et al.: Observational evidence for dissociative shocks in the inner 100 AU of low-mass protostars
are short enough, of the order of a few tens of AU or less.
This is consistent with the dynamical distance, i.e., the distance
a shock with a velocity of 10 km s−1 traverses over 2 years
(∼5 AU). Furthermore, the variability illustrates that the offset components are transient phenomena, and that the driving
agent is not constant over time, i.e., the envelope may have time
to settle between outburst events. The combination of a low
shock velocity (10–15 km s−1 ), dense medium (106 –108 cm−3 ),
UV-driven chemistry and time variability point to a scenario in
which shocks from the protostellar wind impinge on the very
inner dense envelope at angles that are close to perpendicular
to the large-scale outflow. Whether the envelope has enough
time to relax between outburst events or not will await further
observations.
4.4. Tracing winds and shocks in other sources
Fig. 11. Cartoon illustrating the inner few 100 AU of a low-mass protostar (not to scale). Where the wind and UV photons interact directly
with the envelope, an irradiated shock occurs which compresses the envelope by several orders of magnitude over short distances. On the side
of the cavity walls facing the protostar, ions (C+ , CH+ , OH+ ) dominate
emission and cooling. Neutral species (OH, H2 O, CO) dominate further
into the envelope and on larger size scales. The disk hides the innermost
parts of the red-shifted outflow lobe. Typical densities and temperatures
are provided where relevant.
(ii) the hydride column densities, while (iii) not dissipating the
envelope on very short timescales. The reason the fast dissociative shock is successful in reproducing the first two ingredients
is the layered structure where the hottest gas is atomic (dissociated) and as the gas cools, molecules reform. Two alternatives,
which would also reproduce the third ingredient, are possible:
(a) either the dissociating UV photons are not intrinsic to the
shock in which case the shock velocity can be lower, possibly as
low as the wind velocity in Class I/II sources (10–20 km s−1 ), see
Fig. 11 for a schematic of the scenario; or (b) the wind is faster
in Class 0 sources but only interacts with the envelope in very
small regions and over short enough timescales that no irreparable damage is done to the large-scale envelope, i.e., the wind is
anisotropic in time and direction. In the following we argue why
a combination of the two is a likely solution.
Accreting protostars generate copious amounts of UV photons (e.g., Ingleby et al. 2011) and thus provide a natural source
of external UV illumination (Spaans et al. 1995; van Kempen
et al. 2009; Visser et al. 2012; Yıldız et al. 2012). If the UV field
is strong and hard enough, the ion chemistry naturally evolves
in a similar fashion to high-mass star-forming objects along the
outflow cavity walls (e.g., Stäuber et al. 2007; Bruderer et al.
2010a,b; Benz et al. 2010), and the high shock velocity is not
required to dissociate the molecules. Visser et al. (2012) argue
that the cavity density is low enough and that the distance to the
cavity wall from the protostar is short enough that UV photons
reach the cavity wall mostly unattenuated. Depending on the actual shock conditions and the characteristics of the UV field, it
is possible that the combination of a dissociating UV field and
lower shock velocity similar to the wind velocities observed towards Class I and II sources ( ∼ 10–20 km s−1 ) can reproduce
the chemical and physical conditions. However, models of dense
irradiated shocks are required to test this hypothesis further.
The variability observed towards IRAS 4A demonstrates
that heating and cooling in the offset component takes place
on timescales of years. Such a variability is only possible in
J-type shocks (Flower et al. 2003), where the cooling lengths
Most of the neutral species are reformed at a temperature plateau
in the shock of ∼500 K. The heating of this plateau region
is dominated by H2 formation on grains, and CO, H2 O, OH
and other neutral species dominate the cooling. The temperature of this plateau is coincident with that of the hot component
seen in the PACS high-J CO data towards low-mass protostars
(Herczeg et al. 2012; Goicoechea et al. 2012; Manoj et al. 2013;
Karska et al. 2013; Green et al. 2013), and it is possible that the
CO emission arises in the dense post-shock gas. This would imply that almost all protostars contain dense dissociative shocks,
and that the “universality” of this temperature component is due
to the energy release of H2 reformation.
If this interpretation is correct, it provides geometrical constraints on the type of system where the offset component can be
observed. First, the offset component needs to move out of the
plane of the sky, or the radial velocity is zero. If the component
is moving entirely in the plane of the sky, the component is incident with the source velocity and there is no or little offset as
in the case of IRAS 4B. Second, the infall rate needs to be large
enough that the red offset component is shielded by the infalling
gas or a more dense inner envelope or disk. If it is not shielded,
the profile becomes symmetric around the source velocity making identification more difficult. The offset component is only
detected towards the more massive envelopes in the WISH lowmass sample with the higher infall rates and larger disks, which
corroborates this interpretation.
If the infall rate is very low, the corresponding outflow rate is
low and therefore the offset component will be weak because the
shock is weaker. The sources showing the offset component have
the highest mass of hot gas as measured by the very high-J CO
emission observed with PACS (Herczeg et al. 2012; Karska et al.
2013), suggesting that the other reason the offset component is
not observed towards more sources may be a question of the
signal-to-noise ratio. Alternatively, since the variability observed
towards IRAS 4A takes place over timescales of a few years,
the lack of an offset component indicates that the current shock
activity is low.
The offset component is primarily seen in H2 O, and as illustrated above, H2 O is a better tracer of the kinematics than CO
when it comes to hot shocked gas. However, two additional reasons to those mentioned above may play a role as to why an
offset component is not seen towards all sources in H2 O. First,
when the envelope is diluted UV photons may penetrate further
into the envelope and photodissociate H2 O even in the postshock gas, thereby lowering the H2 O column density (Nisini
et al. 2002; Visser et al. 2012; Karska et al. 2013). Second, as
the H2 density becomes lower excitation of H2 O becomes more
A23, page 11 of 13
A&A 557, A23 (2013)
inefficient, and thus the S/N decreases significantly (Nisini et al.
2002; Kristensen et al. 2012). Thus, the reason the offset component is not detected towards all sources is likely a combination
of geometrical effects and low signal from the emitting gas.
5. Summary and conclusions
– Water emission profiles trace components not seen in lowJ CO (J < 10–9) observed from the ground and space-based
telescopes such as Herschel. These components are uniquely
detected in ionised hydrides such as OH+ and CH+ , but also
OH, CO 16–15 and C+ show the same component in absorption and emission. The component is not detected in deep
spectra of HCO+ 6–5 nor CH, thus providing valuable upper
limits on the column densities.
– The component is observed at two epochs towards most
sources and the intensity doubled over a period of two years
in the H2 O 312 –303 transition at 1097 GHz towards IRAS 4A.
The variability points to an origin in gas which is rapidly
heated or cooling.
– H2 O column densities are estimated to be 1016 cm−2 while
the H2 density is 5 × 106 – 5 × 107 cm−3 . The emitting regions are small, ∼100 AU. From the detection of the offset
component in CO 16–15 the CO column density over the
emitting region is estimated to be 1018 cm−2 . The C+ , CH+
and OH+ column densities are measured and upper limits on
CH and HCO+ are provided.
– The inferred CO column density implies that the offset component is identical to the hot CO component
(T ∼ 700–800 K) seen in Herschel-PACS data towards these
low-mass protostars. The temperature is caused by the reformation of H2 in the post-shock gas.
– Several of the light hydrides are observed in absorption
against the continuum, directly pointing to an origin close
to the protostar. The most likely origin of this component
is in dissociative shocks where the dissociation may come
from the UV light from the accreting protostar rather than
the shock itself, thus lowering the required shock velocity.
The shocks possibly arise in the interaction between the protostellar wind and the inner, dense envelope. The UV field
is required to account for the large column densities of the
ionised hydrides, and models of irradiated shocks are necessary to further test and quantify this hypothesis.
These observations highlight the unique capability of, in particular, H2 O as a tracer of dynamical components in protostars even
on the smallest spatial scales. Specifically, the data unequivocally reveal that dissociative shocks in the inner ∼100 AU of
low-mass protostars are common. These observations thus open
a window towards understanding how low-mass protostars energetically interact with their parental material. Future facilities
such as ALMA will undoubtedly shed more light on these energetic processes for a larger number of sources and at the angular
resolution required to spatially resolve these processes.
Acknowledgements. The authors would like to thank the entire WISH team for
many, many stimulating discussions. HIFI has been designed and built by a consortium of institutes and university departments from across Europe, Canada and
the US under the leadership of SRON Netherlands Institute for Space Research,
Groningen, The Netherlands with major contributions from Germany, France
and the US. Consortium members are: Canada: CSA, U. Waterloo; France:
CESR, LAB, LERMA, IRAM; Germany: KOSMA, MPIfR, MPS; Ireland,
NUI Maynooth; Italy: ASI, IFSI-INAF, Arcetri-INAF; Netherlands: SRON,
TUD; Poland: CAMK, CBK; Spain: Observatorio Astronomico Nacional
(IGN), Centro de Astrobiología (CSIC-INTA); Sweden: Chalmers University
A23, page 12 of 13
of Technology – MC2, RSS & GARD, Onsala Space Observatory, Swedish
National Space Board, Stockholm University – Stockholm Observatory;
Switzerland: ETH Zürich, FHNW; USA: Caltech, JPL, NHSC. Astrochemistry
in Leiden is supported by the Netherlands Research School for Astronomy
(NOVA), by a Spinoza grant and grant 614.001.008 from the Netherlands
Organisation for Scientific Research (NWO), and by the European Community’s
Seventh Framework Programme FP7/2007-2013 under grant agreement 238258
(LASSIE).
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L. E. Kristensen et al.: Observational evidence for dissociative shocks in the inner 100 AU of low-mass protostars
Appendix A: Observation IDs
All observations presented in this paper and their respective observation IDs are provided in Table A.1.
Table A.1. Observation IDs of Herschel observations presented and discussed in this paper.
Setting
H2 O 110 –101
H2 O 212 –101
H2 O 111 –000
H2 O 202 –111
H2 O 211 –202
H2 O 312 –303
H2 O 312 –303 a
H2 O 312 –221 /CO 10–9b
OH/CO 16–15b
HCO+ /CHc
CH+d
OH+d
C+d
IRAS 2A
1342202067
1342215966
1342191657
1342191606
1342191748
1342191654
1342215968
1342191701
1342192206
1342203229
1342203180
1342201840
IRAS 3A
1342202066
IRAS 4A
1342202065
1342203951
1342191656
1342191605
1342191749
1342201796
1342249014
1342191721
IRAS 4B
1342202064
1342203952
1342191655
1342191604
1342191750
1342201797
1342249851
1342191722
1342192207
1342203230
1342203181
1342201841
1342202033
1342203228
1342203179
1342201842
SMM1
1342208580
1342207660
1342207379
1342194994
1342194561
1342207378
1342254450
1342207701
1342208574
1342194463
1342207620
1342207656
1342208575
SMM3
1342208579
1342215965
1342207377
1342207658
1342207617
1342207376
1342207699
1342207580
Notes. (a) Reobserved as part of OT2 programme OT2_rvisser_2. (b) Obtained in the same setting. (c) Obtained in the same setting as H18
2 O 110 –101 .
(d)
Observations to be presented in Benz et al. (in prep.).
A23, page 13 of 13